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The Astrophysical Journal, 701:163–175, 2009 August 10
C 2009.
doi:10.1088/0004-637X/701/1/163
The American Astronomical Society. All rights reserved. Printed in the U.S.A.
HIGH-PRECISION C17 O, C18 O, AND C16 O MEASUREMENTS IN YOUNG STELLAR OBJECTS: ANALOGUES
FOR CO SELF-SHIELDING IN THE EARLY SOLAR SYSTEM∗
Rachel L. Smith1 , Klaus M. Pontoppidan2,7 , Edward D. Young1,3 , Mark R. Morris4 , and Ewine F. van Dishoeck5,6
1
Department of Earth and Space Sciences, University of California, Los Angeles, 595 Charles E. Young Drive East, Geology Building, Los Angeles,
CA 90095-1567, USA; [email protected].
2 Division of Geological and Planetary Sciences, California Institute of Technology, Pasadena, CA 91125, USA; [email protected].
3 Institute of Geophysics and Planetary Physics, University of California, Los Angeles, Los Angeles, CA 90095-1567, USA; [email protected].
4 Division of Astronomy and Astrophysics, Department of Physics and Astronomy, University of California, Los Angeles, CA 90095-1547, USA;
[email protected].
5 Leiden Observatory, Leiden University, P.O. Box 9513, NL 2300 RA Leiden, The Netherlands; [email protected].
6 Max Planck Institut für extraterrestrische Physik, Postfach 1312, 85741 Garching, Germany
Received 2008 December 12; accepted 2009 June 8; published 2009 July 21
ABSTRACT
Using very high resolution (λ/Δλ ≈ 95 000) 4.7 μm fundamental and 2.3 μm overtone rovibrational CO absorption
spectra obtained with the Cryogenic Infrared Echelle Spectrograph infrared spectrometer on the Very Large
Telescope (VLT), we report detections of four CO isotopologues—C16 O, 13 CO, C18 O, and the rare species,
C17 O—in the circumstellar environment of two young protostars: VV CrA, a binary T Tauri star in the Corona
Australis molecular cloud, and Reipurth 50, an intermediate-mass FU Ori star in the Orion Molecular Cloud. We
argue that the observed CO absorption lines probe a protoplanetary disk in VV CrA, and a protostellar envelope
in Reipurth 50. All CO line profiles are spectrally resolved, with intrinsic line widths of ≈3–4 km s−1 (FWHM),
permitting direct calculation of CO oxygen isotopologue ratios with 5%–10% accuracy. The rovibrational level
populations for all species can be reproduced by assuming that CO absorption arises in two temperature regimes.
In the higher temperature regime, in which the column densities are best determined, the derived oxygen isotope
ratios in VV CrA are: [C16 O]/[C18 O] = 690 ± 30; [C16 O]/[C17 O] = 2800 ± 300, and [C18 O]/[C17 O]= 4.1 ± 0.4.
For Reipurth 50, we find [C16 O]/[C18 O] = 490 ± 30; [C16 O]/[C17 O] = 2200 ± 150, [C18 O]/[C17 O] = 4.4 ± 0.2.
For both objects, 12 C/13 C are on the order of 100, nearly twice the expected interstellar medium (ISM) ratio.
The derived oxygen abundance ratios for the VV CrA disk show a significant mass-independent deficit of C17 O
and C18 O relative to C16 O compared to ISM baseline abundances. The Reipurth 50 envelope shows no clear
differences in oxygen CO isotopologue ratios compared with the local ISM. A mass-independent fractionation
can be interpreted as being due to selective photodissociation of CO in the disk surface due to self-shielding.
The deficits in C17 O and C18 O in the VV CrA protoplanetary disk are consistent with an analogous origin of
the 16 O variability in the solar system by isotope selective photodissociation, confirmation of which may be
obtained via study of additional sources. The higher fractionation observed for the VV CrA disk compared with the
Reipurth 50 envelope is likely due to a combination of disk geometry, grain growth, and vertical mixing processes.
Key words: astrochemistry – circumstellar matter – ISM: abundances – planetary systems: protoplanetary disks –
stars: individual (VV CrA, Reipurth 50)
α 18/16 = ([18 O]/16 O])/([18 O]/[16 O])◦ . The subscript ◦ refers
to some initial condition. Because mass-dependent fractionation depends on the ratio of partition functions, it is well
known that the relationship is α 17/16 = (α 18/16 )β , where
β = (1/m16 − 1/m17 )/(1/m16 − 1/m18 ) at equilibrium, or
ln(M16 /M17 )/ln(M16 /M18 ) for kinetic processes, and where mi
is the atomic mass for atomic species i, and Mi can be an atomic,
molecular, or reduced mass for isotopologue i depending on the
process (Young et al. 2002). In practice, β ∼ 0.51 (kinetic)
to 0.53 (equilibrium) for oxygen. Isotope variations that obey
this relationship are deemed mass dependent. Mass-independent
variations are those that do not follow this relationship.
Oxygen isotopes in rocks of the solar system exhibit an
anomalous mass-independent distribution that has defied conclusive explanation since its discovery (Clayton et al. 1973).
This anomaly, in the form of a mass-independent correlation
between [16 O]/[18 O] and [16 O]/[17 O] among rocky bodies, is
one of the most pronounced chemical features of the solar system. Here, mass independence refers to relative differences in
[16 O]/[18 O] and [16 O]/[17 O] among primitive rocky objects that
are nearly identical, suggesting changes in [16 O] relative to both
1. INTRODUCTION
The capability to compare isotope ratios of meteorites with
those in protoplanetary systems presents exciting new possibilities for constraining early solar system history. Oxygen is of
particular value in this regard because it occurs in great abundance, comprising ∼50% by mass of most rocky materials and
is, with the [CO]/[H] ratio roughly 10−4 in the gas phase, a significant portion of the gas as well. With the advent of improved
instrumentation, astronomical observations of stable isotope ratios are becoming a potentially powerful new tool for probing
the chemistry of young stellar objects (YSOs), making comparisons with the solar system possible.
Mass-dependent isotope fractionation describes the relationship between the changes in the ratios of three or
more isotopes. In the case of oxygen, one considers the
relationship between the fractionation factors α 17/16 and
α 18/16 , where α 17/16 = ([17 O]/16 O])/([17 O]/[16 O])◦ and
∗ This work is based on observations collected at the European Southern
Observatory Very Large Telescope under program ID 179.C-0151.
7 Hubble Fellow.
163
164
SMITH ET AL.
[17 O] and [18 O], rather than the expected mass-dependent trend
in which relative changes in [16 O]/[17 O] are about half those in
[16 O]/[18 O] on a log scale.
Explanations for the mass-independent distribution of oxygen isotope ratios in the solar system include inheritance of a
nucleosynthesis signal from the interstellar medium (e.g., Galactic chemical evolution), symmetry-based mass-independent isotope fractionation effects such as those involving ozone in
Earth’s stratosphere (Thiemens & Heidenreich 1983), and photochemistry involving isotope-specific photodissociation of CO
(Kitamura & Shimizu 1983; Thiemens & Heidenreich 1983;
Clayton 2002; Yurimoto & Kuramoto 2004; Lyons & Young
2005). Recent results from the Genesis mission establish that
rocks in the solar system are depleted in 16 O relative to the
Sun (McKeegan et al. 2008). The depletion is extreme in aqueously altered primitive meteoritical materials (Sakamoto et al.,
16
16
17
2007), suggesting that [H18
2 O]/[H2 O] and [H2 O]/[H2 O] were
extremely high in the planet-forming region of the early solar
system. Both of these observations lend support to the suggestion that self-shielding by CO oxygen isotopic species during photodissociation of CO is a likely explanation for massindependent fractionation of oxygen isotopes in the solar system
(Young et al. 2008).
Self-shielding by CO refers to the variable shielding of CO
oxygen isotopologues from photodissociation by far ultraviolet
(FUV) radiation in proportion to their abundances (column
densities). Thus, C17 O and C18 O will be more rapidly destroyed
than the much more abundant C16 O. The oxygen liberated
during this process will eventually end up in H2 O, providing an
17
explanation for the overabundance of H18
2 O and H2 O relative
16
to H2 O in the early solar system. Exchange of oxygen isotopes
between newly formed rock and H2 O gas provided a path for
rocks to become enriched in 18 O and 17 O relative to the original
solar oxygen isotopic composition (Yurimoto & Kuramoto
2004; Lyons & Young 2005; Young 2007).
The effect on the survival of oxygen isotopologues is greatest
for column densities of CO in the range 1015 –1018 cm−2 .
Isotope selective photodissociation and self-shielding occur in
molecular clouds (Bally & Langer 1982; van Dishoeck & Black
1988; Sheffer et al. 2002) but their occurrence in circumstellar
disks is yet to be demonstrated conclusively. Models of CO
photodissociation in a disk suggest that isotope selectivity will
be most efficient at the surfaces of the disk regardless of whether
the FUV source is the central star itself or proximal O or B stars
(Lyons & Young 2005; Young 2007). These models predict
that the outer regions of circumstellar disks should exhibit
C17 O and C18 O deficits relative to C16 O as a consequence of
CO self-shielding, providing a test of the models. The precise
degree of isotope selection depends on the amount of radial and
vertical mixing in the disk, as well as on time. Models suggest
C17 O and C18 O deficits relative to C16 O of tens of percent on
timescales of 104 to 105 years (Young 2007). This is more
than sufficient to explain the anomalous distribution of oxygen
isotopes in the solar system. An alternative model (Yurimoto
& Kuramoto 2002; Lee et al. 2008) suggests that signatures of
CO self-shielding in the disk may have been inherited from the
surrounding cloud material, where the excess 17 O and 18 O has
subsequently been incorporated into water ice and transported
to disks; in this case, one expects anomalous deficits in C17 O
and C18 O in the disk and in the envelope surrounding the disk,
regardless of stage of evolution.
In this work, we present column densities for the CO isotopologues 12 C16 O (C16 O), 12 C18 O (C18 O), 12 C17 O (C17 O), and
Vol. 701
13 16
C O (13 CO) in two YSOs that comprise a first test of the
importance of CO self-shielding in the early solar system. We
emphasize that quantifying both [16 O]/[18 O] and [16 O]/[17 O]
is particularly important in this study because only the combination of the two ratios can distinguish conclusively between
photochemical self-shielding and mass-dependent fractionation.
Infrared absorption spectroscopy has been used previously to
measure the ratios of the 12 CO, 13 CO, and C18 O isotopologues
in YSOs. Infrared observations of one circumstellar gaseous
disk provide a hint of C16 O overabundance relative to C18 O:
Brittain et al. (2005) presented high-resolution infrared spectra
of the embedded, low-mass pre-main sequence star HL Tau,
showing a ratio of column densities for C16 O and C18 O,
N (C16 O)/N (C18 O), of 800 ± 200 (2σ ). Typical [16 O]/[18 O]
in the local interstellar medium (ISM) is 557 ± 30, (Wilson
1999), similar to the solar system value of 499. The apparent
overabundance of C16 O in the HL Tau disk could be the result
of isotope-selective photodissociation resulting from the HL
Tau UV field, although the uncertainties do not yet allow
one to completely rule out standard ISM isotopologue ratios.
Interpretation of these data are complicated by the fact that
the line of sight toward HL Tau samples mostly the envelope
rather than the disk. In addition, without the complementary
data for C17 O, it is not possible to differentiate between massindependent versus mass-dependent fractionation mechanisms;
at temperatures that exist in the outer regions of disks, mass
fractionation could be 20% or more (Young 2007), well within
the range for HL Tau.
Searching for oxygen isotope fractionation in protoplanetary
disks requires an observational method that can measure the
abundances of the four most common CO isotopologues to
relative accuracies better than 10%. Isotopologue ratios measured using rotational emission lines depend on excitation and
radiative transfer models with significant uncertainties due to
non-thermalized level populations and differences in beam sizes
arising from differences in the frequencies of the various CO
isotopologues. Further, rotational emission is currently limited
to the lowest rotational levels (J 3), and is therefore tracing
low-temperature gas only. An alternative that eliminates most
of these concerns is to use absorption lines in which transitions
from different isotopologues are known to probe nearly identical pencil beam lines of sight. The fundamental rovibrational
band at 4.7 μm and the overtone at 2.3 μm of CO provide such
an absorption-line tracer. We show that very high-resolution
infrared spectroscopy can be used to derive accurate column
densities of all three isotopologues of 12 CO. The main challenge is to locate sources with favorable geometries providing
column densities high enough to allow the detection of C17 O,
but not so high as to extinguish the infrared source. One type of
favorable absorption-line geometry is that of a disk viewed at
high inclination angles to allow for absorption in the outer disk
of infrared continuum emission originating in the inner disk. In
this paper, we take advantage of a new type of geometry, namely
that of a binary star in which one star is being eclipsed by the
disk of the other star. We compare one such binary T Tauri star
with an embedded YSO in which the absorption may in part be
attributed to a remnant envelope.
2. SOURCE SELECTION
The first absorption-line source is the northern component of
the binary T Tauri star VV CrA (J2000: R.A. = 19 03 06.8,
decl. = −37 12 54.6); the southern component does not show
No. 1, 2009
CO ISOTOPOLOGUE RATIOS IN PROTOSTELLAR DISKS AND ENVELOPES
165
Cold component (midplane)
Warm component (surface layer)
270 AU
VV CrA North
VV CrA South
100
Case A
AU
VV CrA North
Cold component
Warm component
Case B
Figure 1. Schematic cartoon illustrating the two possible geometries for the VV CrA binary system. Case A illustrates the scenario whereby the disk of the secondary
star is eclipsed by the outer disk of the primary; Case B shows the less probable, single inclined disk geometry.
CO in absorption. VV CrA N is part of the class of “infrared
companion” (IRC) sources, which are highly extincted companions to optically visible T Tauri stars with projected separations
of a few hundred AU or less (Koresko et al. 1997). Infrared
companions typically show silicate absorption features, revealing that the absorbing material is localized. One interpretation
is that the infrared companions have very compact, dusty envelopes. Here, we suggest a possible alternative explanation:
that at least some of the infrared companions are regular T Tauri
stars that are being eclipsed by the outer disk of the primary
star of the binary (e.g., Hogerheijde et al. 1997). The only strict
requirement for this scenario is that the disks of the two companions be non-coplanar, a geometry that gains support from
the fact that non-coplanarity is a known property of some binary T Tauri stars (e.g., Koresko 1998; Stapelfeldt et al. 2003).
The geometrical situation we envision is depicted schematically
in Figure 1. While the available data do not allow a distinction
between the two-disk scenario (Case A) and one in which the
gas is located in the outer parts of a single inclined disk (Case
B), the former scenario is geometrically more likely. In either
case, we are observing gas in the outer disk, and thus our analysis of the oxygen isotope ratios via CO is not predicated on the
differentiation between the two geometries, although we note
that Case B will minimize systematic differences between lines
of sight at 2.3 and 4.7 μm.
The projected separation of the VV CrA components is 2. 1,
corresponding to 273 AU at the distance of 130 pc to the Corona
Australis star-forming region (Casey et al. 1998; Neuhäuser
& Forbrich 2008). The VV CrA system is highly variable, as
indicated by a K-band flux of −2.3 mag in 1987 (Chelli et al.
1995), 0.3–0.0 mag in 1993/95 (Ageorges et al. 1997; Ghez et al.
1997) and 1.3 mag at the time of our observation in 2007. Due
to the fact that the southern component has remained relatively
constant at K ∼6 mag, this corresponds to the IRC fading by
∼3 mag in K in the past 20 years. Given the geometry we
propose in Figure 1 (Case A), this variability could be a result
either of intrinsic variability of the continuum source, or to
column density variations in the orbiting disk of the companion
that move across our line of sight to the continuum source. It
is unlikely to be due to temperature or composition changes on
such short timescales.
The second source of this study is Reipurth 50 [RE 50]
(IRAS 05380-0728) (J2000: R.A. = 05 40 27.7, decl. = −07
27 28), a luminous (250 L ), embedded YSO in the Orion
star-forming cloud at a distance of 470 pc. It illuminates a
large, bright reflection nebula that appeared sometime between
1960 and 1970 (Reipurth & Bally 1986; Reipurth & Aspin
1997). Strom & Strom (1993) found that Reipurth 50 is an FU
Ori type YSO, and it has been shown to be surrounded by a
compact 0.1 M envelope (Sandell & Weintraub 2001), placing
it in stage I (Robitaille et al. 2006) of the YSO evolutionary
sequence. It is likely that the absorbing gas is located in this
envelope, rather than in a disk. Thus, Reipurth 50 and VV CrA
may provide a comparison between disk and envelope material,
although larger samples of objects are required to make any firm
conclusions.
3. OBSERVATIONS
The observations of VV CrA and Reipurth 50 were taken as
part of a large program to observe about 100 YSOs and protoplanetary disks with the newly implemented Cryogenic Infrared
Echelle Spectrograph (CRIRES) at the Very Large Telescope
(VLT) in Chile. CRIRES is an AO-assisted spectrometer that
166
SMITH ET AL.
Table 1
Journal of Observations
Source/Obs. Date
(UT)
Spectral Range
(μm)
Integrationa
(minutes)
Signal/Noiseb
VV CrA/2007-08-31
VV CrA/2007-08-31
Reipurth 50/2007-10-11
Reipurth 50/2007-10-17
4.701 − 4.815
2.338 − 2.396
4.645 − 4.901
2.338 − 2.397
32
23
10–20
20
260
210
290
40
Notes.
a Integration is variable with wavelength, so values are approximate.
b Median values.
operates at very high resolving powers (for infrared spectroscopy) of λ/Δλ ≈ 100 000. The spectra of these sources
allow for accurate determination of the CO column densities
and isotopologue ratios. Table 1 summarizes the observing parameters for VV CrA and Reipurth 50.
The spectrum of VV CrA was observed on 2007 August
31 in both the M (4.7 μm) and K (2.3 μm) bands. Reipurth
50 was observed on 2007 October 11 and 17, respectively,
and the observations from both dates were combined. CRIRES
M-band spectra of VV CrA obtained in 2007 April showed
variations in the depths of the CO lines of ∼10% relative
to the August data, illustrating the importance of obtaining
nearly concurrent K- and M-band spectra. The spectra of the
v = (1–0) fundamental rovibrational bands for C17 O, C18 O,
and 13 CO and the v = (2–0) overtone band of C16 O for VV
CrA were obtained using the 0. 2 slit, resulting in a resolving
power of R = λ/Δλ ≈ 95 000 (corresponding to 3.16 km
s−1 ), as measured on unresolved telluric lines. Several settings
were obtained for each source to probe a range of rotational
levels for all the isotopologues, spanning J = 0 to at least
J = 9. For 12 CO, the v = (2–0) transitions were used because
they have much smaller line strengths (by a factor of ∼130)
than the fundamental, and are thus expected to be optically thin
(or nearly so).
The data were reduced using standard procedures, including
flat-field correction, adjustments to account for detector nonlinearity and linearization of the spectra in both the dispersionand cross-dispersion directions. The spectra were wavelength
calibrated using the telluric absorption lines referenced to
an atmospheric model spectrum and transformed to the local
standard of rest frame. Relative flux calibration was carried out
by dividing the target sources by standard spectra of the stars
HR 7236 (B9V) and HR 1666 (A3III). For more information on
the CRIRES data processing, see Pontoppidan et al. (2008).
4. RESULTS AND DATA ANALYSIS
The spectra of the fundamental and overtone CO bands are
shown in Figures 2 and 3, respectively, illustrating the forest
of narrow absorption lines due to CO gas. While we are only
interested in the narrow CO absorption lines, it is noteworthy
that VV CrA also shows broad, complex emission from CO,
probably originating in the innermost (R < 1 AU) regions of the
disk. Reipurth 50 does not exhibit an emission component, but
the known absorption band from solid CO centered on 4.67 μm
can be seen. Given that the respective emission and ice features
in these sources are much broader than the absorption lines,
their influence on derived gas CO column densities is minimal.
In addition, the velocity shift of the center of the broad emission
lines with respect to the narrow absorption lines are consistent
with, and supportive of, our picture in which the absorption takes
Vol. 701
place in a different structure than the disk around the continuum
source.
Multiple lines from the four most abundant isotopologues
of CO are detected in both sources, including the rare species
C17 O (Figure 4). C17 O has been observed in the ultraviolet
toward the X Persei molecular cloud (Sheffer et al. 2002) and
the P (1) C17 O rovibrational line has been recently reported in
the T Tauri star SR21 in infrared emission (Pontoppidan et al.
2008). While the latter observation would potentially allow
a window into measuring isotope ratios in the inner <1 AU
regions of disks and unambiguous determination of the location
of the gas, radiative transfer modeling for emission lines is
more complicated than for pencil beam absorption lines. We
emphasize that our observations of all four CO isotopologues in
infrared absorption probe a single sight line, therefore obviating
the need for complex modeling required for emission lines.
4.1. Derivation of Column Densities
In order to determine the total column density for each
isotopologue, the column density of CO in each rotational state
must be measured; the total column density of CO can then be
calculated by summing the measured column densities in the
observed rovibrational levels and assuming that the remainder
are populated according to a Boltzmann distribution. If lines
are spectrally unresolved, accurate abundance derivations are
complicated by the fact that it is difficult to determine optical
depths, necessitating a curve-of-growth analysis. Given the high
spectral resolution of CRIRES, the data reported in this study are
not complicated by these factors, and we are able to distinguish
intrinsic line broadening from that of the instrument. Figure 4
is a representative section of the fundamental CO band in
Reipurth 50, illustrating the quality of the data. The CO lines
were confirmed to be spectrally resolved by comparing them
to adjacent unresolved telluric lines; the total CO line widths
are slightly, but significantly, broader than the instrumental
resolution of 3.2 km s−1 exhibited by the telluric lines. Column
densities were obtained from the optical depths at line center
(τ◦ ) derived from the best-fit profiles to each line. Tables 2
and 3 show the optical depths at line center and Doppler
shifts (VLSR ) for each observed line used in the derivation of
isotopologue column densities for VV CrA and Reipurth 50,
respectively.
The fundamental v = (1–0) C16 O lines are optically thick.
Therefore, the C16 O column densities were measured using
v = (2–0) transitions in order to use only lines that are
as optically thin as possible. All line fits are consistent in
having a single, constant width, within error, as expected for
optically thin lines. That width is substantially larger than the
thermal linewidth (see below), and so is likely due to some
combination of local turbulent motions and radial velocity
gradients transverse to the absorbing column.
We fit the frequency-dependent flux, Iν , for any given
transition using
Iν = Ic exp(−s φν NvJ ),
(1)
where Ic is the flux of the continuum and φν is the intrinsic
line profile (normalized to unity). The intrinsic line profile is
assumed to be Gaussian with a full width at half maximum
(FWHM) = γintrinsic . NvJ is the column density of the lower
level of the transition, having rotational and vibrational quantum
numbers J and v (v = 0 for all observations presented here), s
is the integrated cross section, and τ◦ ≡ sNvJ φ◦ .
No. 1, 2009
CO ISOTOPOLOGUE RATIOS IN PROTOSTELLAR DISKS AND ENVELOPES
VV CrA
υ = (1 − 0)
1.4
16
C O: P(6) - P(15)
16
C O: R(7) - P(4)
12 18
C O: R(9) - P(3)
12 17
C O: R(0) - P(9)
13
CO emission
1.2
Normalized Flux
12
Lines detected:
167
1.0
0.8
12
0.6
17
C O
13
0.4
12
18
16
C O
C O
0.2
12
16
C O
4.70
4.72
4.74
4.76
4.78
4.80
Wavelength [μm]
1.4
Lines detected:
Reipurth 50
υ = (1 − 0)
Normalized Flux
1.2
12
1.0
12
16
C O: R(0) - P(24)
16
C O: R(15) - P(13)
12 18
C O: R(16) - P(12)
12 17
C O: R(9) - P(9)
13
16
C O ice
0.8
0.6
12
0.4
18
C O
12
13
0.2
12
16
C O
4.66
17
C O
16
C O
4.68
4.70
4.72
4.74
4.76
4.78
Wavelength [μm]
Figure 2. Infrared absorption spectra of the CO fundamental rovibrational bands toward VV CrA and Reipurth 50. The fully observed spectral range is represented
for VV CrA; for Reipurth 50, the portion with absorption lines used for the analysis is shown. Absorption lines are due to various CO isotopologues, representatives
of which are indicated. Note that, in addition to the narrow absorption lines, VV CrA shows complex, low-level, broad-lined emission from hot CO gas, presumably
located in the inner disk. The broad 12 C16 O ice feature in Reipurth 50 is also marked.
The integrated cross section (cm2 s−1 ) (uncorrected for stimulated emission) is
s=
π e2
fv J ←vJ ,
me c
(2)
(Spitzer 1978), where fv J ←vJ is the oscillator strength for
absorption from level υ, J to level υ ,J , and is related to the
Einstein A coefficient by
Av J →vJ =
8ν 2 π 2 e2 gJ
fv J ←vJ
me c 3 gJ (3)
Einstein As were taken from the HITRAN database (Rothman
2005), which uses the values of Goorvitch & Chackerian (1994a, 1994b). Ayres et al. (2006) compared different
sets of molecular parameters for the CO fundamental band
and found our adopted values to be consistent with other
works.
Because the intrinsic line widths are close to the instrumental
resolution, the Iν given by Equation (1) is convolved with the
instrumental line profile modeled by a Gaussian with FWHM =
3.2 km s−1 prior to fitting the line spectrum. We adopted the
median intrinsic width of the C18 O lines as the best estimate of
the line width and fixed this parameter in the fits to individual
lines. The measured mean FWHMs of the absorption profiles are
4.5 km s−1 for Reipurth 50 and 5.1 km s−1 for VV CrA, which,
when deconvolved with the instrumental profile, yield γintrinsic =
3.3 and 4.0 km s−1 , respectively. The scatter of measured line
widths suggests that the uncertainty on the median width is
10%–20%. The value of γintrinsic
√is related to the Doppler velocity
dispersion, b, by γintrinsic = 2 ln 2 b. The baseline intensities,
Ic , were fit simultaneously with the line profiles using first-order
polynomial fits to the local continuum. Sample fits are illustrated
in Figure 5.
From these fits we obtained rotational diagrams for the respective sources, shown in Figure 6. These plots show the sublevel column densities, NJ /(2J + 1), derived from the fits to
the absorption lines, versus the energy of the state J. A linear
trend is expected for a single-temperature gas. It is clear from
the different slopes between the low-J transitions (J 3) and
higher-J transitions that there is a range of different gas temperatures along the line of sight toward both sources. This is to
be expected for essentially all geometries in which the gas is
heated by the central star. While assuming a continuous distribution of temperatures might be most realistic, we found that
a distribution consisting of only two temperatures is sufficient
168
SMITH ET AL.
Vol. 701
VV CrA
Normalized Flux
1.2
12
16
C O
υ = (2 − 0)
1.0
P(15)
P(14)
P(13)
P(10)
0.8
P(1)
R(3) R(2)
0.6
P(3)
P(2)
R(0)
P(4)
P(6)
P(5)
P(7)
P(8)
P(11)
P(12)
P(9)
R(1)
0.4
2.34
2.36
2.35
2.37
2.38
Wavelength [μm]
Reipurth 50
Normalized Flux
1.2
12
16
C O
υ = (2 − 0)
1.0
P(11)
0.8
P(9)
0.6
R(0)
P(7)
P(10)
P(8)
P(5)
P(3)
P(2)
R(1)
0.4
P(6)
P(1)
P(12)
P(4)
R(3) R(2)
2.34
2.36
2.35
2.37
2.38
Wavelength [μm]
Figure 3. Infrared absorption spectra of the CO overtone rovibrational band toward VV CrA [P(16)–P(17) not shown] and Reipurth 50. Detected 12 C16 O lines are
marked.
1.4
Reipurth 50
= (1 0)
Normalized Flux
1.2
1.0
12
17
C O
P(2)
0.8
12
18
C O
R(4)
0.6
0.4
13
0.2
16
C O
R(3)
4.739
4.740
4.741
4.742
Wavelength [μm]
Figure 4. Representative magnification of CO rovibrational lines for rare isotopologues in Reipurth 50.
to capture the distribution of points in the excitation diagrams
for both sources. In the absence of any explicit information
on the actual temperature distributions, we proceeded with a
simultaneous two-temperature fit to the data for each isotopologue (Figure 6); from these models, individual low and high
temperatures were derived. For both objects, the two temper-
atures consisted of a cold component of T ∼ 10–20 K and a
warm component of T ∼ 150–250 K. In the case of VV CrA,
it is worth noting that the temperature of the warm component
is roughly consistent with that of the uppermost surface layers at a few 100 AU in detailed disk heating models (Kamp &
Dullemond 2004; Jonkheid et al. 2004). The column density of
No. 1, 2009
CO ISOTOPOLOGUE RATIOS IN PROTOSTELLAR DISKS AND ENVELOPES
Table 2
Optical Depths and Doppler Shifts for Observed CO Isotopologue
Rovibrational Lines used in Deriving Column Densities in VV CrAa
VV CrA
Isotopologue
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
Line ID
Optical Depth
(τ◦ )
Doppler Shift
(VLSR , km s−1 )
(2, 0) R(3)
(2, 0) R(2)
(2, 0) R(1)
(2, 0) R(0)
(2, 0) P (1)
(2, 0) P (2)
(2, 0) P (3)
(2, 0) P (4)
(2, 0) P (5)
(2, 0) P (6)
(2, 0) P (7)
(2, 0) P (8)
(2, 0) P (9)
(2, 0) P (10)
(2, 0) P (11)
(2, 0) P (12)
(2, 0) P (13)
(2, 0) P (14)
(2, 0) P (15)
(2, 0) P (16)
(2, 0) P (17)
(1, 0) R(7)
(1, 0) R(6)
(1, 0) R(5)
(1, 0) R(4)
(1, 0) R(3)
(1, 0) R(2)
(1, 0) R(0)
(1, 0) P (1)
(1, 0) P (2)
(1, 0) P (4)
(1, 0) R(9)
(1, 0) R(7)
(1, 0) R(6)
(1, 0) R(5)
(1, 0) R(4)
(1, 0) R(3)
(1, 0) R(2)
(1, 0) R(1)
(1, 0) R(0)
(1, 0) P (1)
(1, 0) P (2)
(1, 0) P (3)
(1, 0) R(0)
(1, 0) P (1)
(1, 0) P (2)
(1, 0) P (3)
(1, 0) P (5)
(1, 0) P (6)
(1, 0) P (8)
(1, 0) P (9)
0.475 ± 0.005
0.535 ± 0.006
0.581 ± 0.006
0.371 ± 0.005
0.309 ± 0.005
0.397 ± 0.004
0.337 ± 0.005
0.319 ± 0.005
0.371 ± 0.005
0.381 ± 0.004
0.360 ± 0.005
0.310 ± 0.005
0.287 ± 0.005
0.255 ± 0.009
0.231 ± 0.005
0.206 ± 0.004
0.154 ± 0.005
0.136 ± 0.005
0.088 ± 0.007
0.093 ± 0.007
0.064 ± 0.005
1.038 ± 0.010
1.381 ± 0.006
0.880 ± 0.007
1.159 ± 0.007
1.104 ± 0.008
1.450 ± 0.008
0.921 ± 0.005
0.611 ± 0.006
1.025 ± 0.010
1.091 ± 0.007
0.124 ± 0.003
0.180 ± 0.004
0.112 ± 0.004
0.179 ± 0.004
0.186 ± 0.003
0.164 ± 0.006
0.136 ± 0.003
0.180 ± 0.005
0.141 ± 0.004
0.102 ± 0.004
0.103 ± 0.004
0.133 ± 0.006
0.028 ± 0.004
0.027 ± 0.003
0.039 ± 0.004
0.024 ± 0.004
0.045 ± 0.003
0.034 ± 0.005
0.022 ± 0.003
0.034 ± 0.004
4.16 ± 0.03
4.77 ± 0.02
5.08 ± 0.02
5.61 ± 0.03
5.41 ± 0.04
5.01 ± 0.03
4.21 ± 0.03
3.57 ± 0.04
3.55 ± 0.04
3.67 ± 0.03
3.62 ± 0.03
3.53 ± 0.04
3.43 ± 0.04
3.46 ± 0.08
3.78 ± 0.05
4.08 ± 0.05
3.99 ± 0.08
4.78 ± 0.09
3.94 ± 0.16
3.72 ± 0.17
4.28 ± 0.21
3.17 ± 0.02
3.25 ± 0.01
3.08 ± 0.01
3.38 ± 0.01
3.30 ± 0.01
4.00 ± 0.01
5.09 ± 0.01
4.23 ± 0.02
4.40 ± 0.02
3.23 ± 0.01
2.75 ± 0.06
3.32 ± 0.07
2.79 ± 0.08
3.38 ± 0.05
3.39 ± 0.04
3.20 ± 0.09
3.81 ± 0.06
5.02 ± 0.06
5.79 ± 0.07
5.21 ± 0.09
3.53 ± 0.09
3.20 ± 0.11
5.72 ± 0.31
5.72 ± 0.26
4.72 ± 0.25
3.86 ± 0.40
2.82 ± 0.18
3.83 ± 0.36
3.10 ± 0.38
4.27 ± 0.32
Note.
a Uncertainties are 1σ derived from the line fits.
the warm component is significantly larger than expected from
these models perpendicular to the midplane, but may be explained if small dust grains have been depleted from the upper
disk layers with a concomitant drop in opacity. A non-vertical,
slanting line of sight through the disk can also enhance the
total column of warm gas. Further modeling of these specific
sources is needed to determine if the observed temperatures and
169
column densities are quantitatively consistent with flared disk
models.
In both VV CrA and Reipurth 50, the derived low temperatures are consistent with the outer disk near the midplane or,
in the case of Reipurth 50, possibly the surrounding envelope.
The high-temperature contributions to the measured absorption
strengths of lines dominated by the cold component are considerable (Figure 6).
The uncertainties in individual line depths are dominated by
systematics, and these are taken into account in the final errors
for the isotopologue ratios. The greatest source of systematic
uncertainty is likely to be residuals from telluric lines that
were not completely removed. This source of error is difficult
to quantify, and varies strongly for different lines. To account
for this, we added 5% uncertainties in quadrature to each line
strength for most of the isotopologues in both temperature
regimes; exceptions to the 5% uncertainty were accounted for
in VV CrA for the low temperature 12 C17 O, low temperature
12 18
C O, and high temperature 13 C16 O lines, for which no
systematic uncertainty, 10% and 20% systematic uncertainty
were included, respectively. These variations resulted from
variable optimizations in fits to obtain reduced χ 2 values near
unity; χ 2 ellipses are shown in the temperature-column density
plane (Figures 7 and 8). A few lines are clearly outliers relative
to a linear fit in the rotational plot by many standard deviations;
these lines, shown as faded points in the fits in Figure 6, were
excluded from the final fits. Inclusion of the estimates for
systematic errors means the uncertainties cited here are likely
to be conservative.
4.2. Isotopic Ratios and Uncertainties
Derived excitation temperatures and isotopologue ratios are
listed in Tables 4 and 5, respectively. There is a significant
difference in excitation temperatures among the different isotopologues as evidenced by the variable slopes in Figure 6 for
each temperature component, apparently violating the assumption of local thermodynamic equilibrium. For example, the excitation temperature for the high-temperature 13 CO component
in Reipurth 50 is greater than that of the other isotopologues
(Figure 8). A similar effect is seen in the warm VV CrA 12 CO
lines (Figure 7). However, the apparently higher temperature of
12
CO could be due to line trapping in optically thick rovibrational lines, in which energy escapes from the rarer isotopologues more readily than from 12 CO due to high optical depth
in the latter. In this case, the derived isotopologue ratios should
not be affected.
Because the 12 CO lines are observed in a different wavelength
region (2.3 μm) than the other isotopologues (4.7 μm), it is
possible that the source location and size are not exactly the same
due to the greater scattering of the K band relative to the M band.
This could cause significant differences in the column densities
that are probed. If the absorbing gas is very close to an infrared
emitting source that changes in size or position with wavelength,
it is still possible that we are probing different volumes and
perhaps path lengths. However, in a spherical envelope, material
with temperatures of 150–200 K is located at distances of ∼100–
200 AU from a 250 L source, compared to a size of only a
few AU for the infrared source. Hence, this is potentially the
largest source of error for both an embedded source with a large
reflection nebula like Reipurth 50, as well as a single inclined
disk geometry as illustrated by Case B (Figure 1) for VV CrA.
The more probable two-disk scenario for VV CrA (Case A) is
not likely to suffer from optical path differences. With CRIRES,
170
SMITH ET AL.
Vol. 701
Table 3
Optical Depths and Doppler Shifts for Observed CO Isotopologue
Rovibrational Lines Used in Deriving Column Densities in Reipurth 50a
Table 3
(Continued)
Reipurth 50
Reipurth 50
Isotopologue
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
12 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
13 C16 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C18 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
12 C17 O
Isotopologue
Line ID
Optical Depth
(τ◦ )
Doppler Shift
(VLSR , km s−1 )
(2, 0) R(3)
(2, 0) R(2)
(2, 0) R(1)
(2, 0) R(0)
(2, 0) P (1)
(2, 0) P (2)
(2, 0) P (3)
(2, 0) P (4)
(2, 0) P (5)
(2, 0) P (6)
(2, 0) P (7)
(2, 0) P (8)
(2, 0) P (9)
(2, 0) P (11)
(2, 0) P (12)
(1, 0) R(15)
(1, 0) R(13)
(1, 0) R(12)
(1, 0) R(11)
(1, 0) R(10)
(1, 0) R(7)
(1, 0) R(6)
(1, 0) R(5)
(1, 0) R(4)
(1, 0) R(3)
(1, 0) R(2)
(1, 0) R(0)
(1, 0) P (1)
(1, 0) P (4)
(1, 0) P (9)
(1, 0) P (10)
(1, 0) P (11)
(1, 0) P (13)
(1, 0) R(15)
(1, 0) R(14)
(1, 0) R(13)
(1, 0) R(12)
(1, 0) R(11)
(1, 0) R(10)
(1, 0) R(9)
(1, 0) R(7)
(1, 0) R(6)
(1, 0) R(5)
(1, 0) R(4)
(1, 0) R(2)
(1, 0) R(1)
(1, 0) R(0)
(1, 0) P (2)
(1, 0) P (3)
(1, 0) P (8)
(1, 0) P (9)
(1, 0) P (10)
(1, 0) R(9)
(1, 0) R(8)
(1, 0) R(7)
(1, 0) R(6)
(1, 0) R(5)
(1, 0) R(4)
(1, 0) R(2)
(1, 0) R(0)
(1, 0) P (2)
(1, 0) P (3)
1.013 ± 0.05
1.035 ± 0.04
0.847 ± 0.04
0.611 ± 0.03
0.619 ± 0.04
0.775 ± 0.05
0.799 ± 0.04
0.796 ± 0.04
0.752 ± 0.05
0.658 ± 0.03
0.609 ± 0.03
0.486 ± 0.04
0.291 ± 0.03
0.220 ± 0.03
0.189 ± 0.03
0.314 ± 0.004
0.558 ± 0.006
0.746 ± 0.010
0.881 ± 0.009
1.044 ± 0.008
1.451 ± 0.012
1.510 ± 0.006
1.877 ± 0.007
1.955 ± 0.007
2.099 ± 0.008
2.449 ± 0.011
1.840 ± 0.006
1.612 ± 0.005
1.978 ± 0.012
1.244 ± 0.008
1.087 ± 0.008
0.897 ± 0.009
0.611 ± 0.007
0.025 ± 0.003
0.051 ± 0.004
0.084 ± 0.006
0.101 ± 0.005
0.142 ± 0.004
0.196 ± 0.003
0.239 ± 0.004
0.372 ± 0.003
0.450 ± 0.003
0.485 ± 0.003
0.457 ± 0.003
0.651 ± 0.003
0.568 ± 0.004
0.348 ± 0.002
0.456 ± 0.004
0.516 ± 0.005
0.312 ± 0.004
0.238 ± 0.006
0.193 ± 0.004
0.049 ± 0.003
0.068 ± 0.003
0.082 ± 0.004
0.097 ± 0.006
0.114 ± 0.006
0.128 ± 0.004
0.132 ± 0.003
0.082 ± 0.003
0.115 ± 0.003
0.105 ± 0.002
5.31 ± 0.09
4.91 ± 0.08
4.94 ± 0.10
4.35 ± 0.11
4.85 ± 0.12
4.69 ± 0.13
5.08 ± 0.09
5.36 ± 0.10
4.89 ± 0.13
4.90 ± 0.11
5.67 ± 0.12
5.79 ± 0.15
5.81 ± 0.23
5.26 ± 0.30
5.58 ± 0.33
5.59 ± 0.03
5.59 ± 0.02
5.66 ± 0.03
5.84 ± 0.02
5.63 ± 0.01
5.43 ± 0.01
5.44 ± 0.01
5.13 ± 0.00
5.26 ± 0.00
5.26 ± 0.00
4.94 ± 0.01
4.74 ± 0.00
5.17 ± 0.00
5.08 ± 0.01
3.61 ± 0.01
5.56 ± 0.01
6.09 ± 0.02
4.35 ± 0.02
5.56 ± 0.30
5.67 ± 0.19
6.02 ± 0.18
6.05 ± 0.12
5.66 ± 0.06
5.44 ± 0.04
5.40 ± 0.04
5.55 ± 0.02
5.52 ± 0.01
5.78 ± 0.01
5.21 ± 0.02
5.08 ± 0.01
5.13 ± 0.01
5.15 ± 0.01
5.15 ± 0.02
5.21 ± 0.02
3.90 ± 0.03
5.57 ± 0.06
5.98 ± 0.05
5.36 ± 0.15
5.29 ± 0.10
5.68 ± 0.11
5.78 ± 0.13
5.87 ± 0.11
5.27 ± 0.07
4.80 ± 0.05
5.10 ± 0.08
5.08 ± 0.06
5.25 ± 0.05
12 C17 O
12 C17 O
12 C17 O
12 C17 O
Line ID
Optical Depth
(τ◦ )
Doppler Shift
(VLSR , km s−1 )
(1, 0) P (4)
(1, 0) P (5)
(1, 0) P (8)
(1, 0) P (9)
0.097 ± 0.003
0.106 ± 0.003
0.054 ± 0.005
0.041 ± 0.003
5.36 ± 0.06
5.48 ± 0.06
6.07 ± 0.19
5.47 ± 0.17
Note.
a Uncertainties are 1σ derived from the line fits.
the pointing is the same for all spectral settings, with both
K- and M-band observations tracking sources in the K band.
We therefore believe that there are no significant differences in
the absorbing volumes for the K and M bands for Reipurth 50
and VV CrA.
The low temperatures derived from our model are low enough
that they could well correspond to the temperature of the
surrounding molecular cloud for Reipurth 50, or the outer disk in
Case A or B for VV CrA, perhaps weakly heated by the nearby
protostar. However, due to the significant contribution of the
high-temperature component on the low-temperature lines, we
believe that our low-temperature ratios contain a much larger
degree of uncertainty than we are able to precisely quantify.
We therefore do not believe these numbers are meaningful.
In discussing isotope ratios, we therefore refer to the ratios
obtained from the high-temperature lines only.
5. DISCUSSION
Both VV CrA and Reipurth 50 should have [C16 O]/[C18 O]
and [C16 O]/[C17 O] values similar to the local ISM in the
absence of extensive isotope partitioning by photochemistry
(self-shielding) or mass-dependent kinetics. The oxygen isotope
ratios for the ISM as a function of distance from the Galactic
center (Galactocentric radius, RGC ) can be estimated from the
4 kpc and 8 kpc values reported by Wilson (1999), yielding for
[16 O]/[18 O],
[16 O]/[18 O]RGC = (57.5 ± 10) RGC + 97,
(4)
where the 1σ uncertainty in the slope reflects the uncertainty in
the means (standard errors) reported in Table 4 of (Wilson 1999).
For [16 O]/[17 O] we can combine Equation (4) with values for
[18 O]/[17 O] in the ISM. These estimates vary from 3.5 (Wilson
1999) to 4.1 (Wouterloot et al. 2008).
Based on Equation (4) and [18 O]/[17 O] = 4.1 for the ISM,
one should expect CO surrounding VV CrA (RGC = 7.8 kpc)
to have [C16 O]/[C18 O] of 550 ± 90 and [C16 O]/[C17 O] of
2300 ± 250. The measured values of 690 ± 30 (1σ ) and
2800 ± 300 for VV CrA are greater. Conversely, the expected
[C16 O]/[C18 O] and [C16 O]/[C17 O] values for Reipurth 50
(RGC = 8.4 kpc), 580 ± 90 (1σ ) and 2400 ± 270, respectively,
are indistinguishable from the measured values of 490 ± 30
(1σ ) and 2200 ± 150 (Table 5).
Results for the oxygen isotopologues in the high-temperature
regime are summarized in the three-isotope plot shown in
Figure 9, where differences in [16 O]/[18 O] and [16 O]/[17 O]
from local ISM values are shown in per mil using the linearized form of the delta notation commonly used in isotope
cosmochemistry. Here, δ 18 O = 103 ln(18 Ri /18 RLocal ISM ) and
CO ISOTOPOLOGUE RATIOS IN PROTOSTELLAR DISKS AND ENVELOPES
171
Optical depth
No. 1, 2009
Velocity [km/s]
Figure 5. Representative fits of CO isotopologues for VV CrA and Reipurth 50. Fits to C16 O lines are to the overtone band. All other fits are to the fundamental
spectra. The fits are shown here normalized to the linear best-fit baseline. Error bars are 1σ /velocity channel. Statistical errors do not include telluric residuals (see
text).
Ri = [C18 O]/[C16 O]. Error ellipses in the three-isotope plot
were calculated using a Monte Carlo calculation that propagates uncertainties in column densities to the final delta values.
The three-isotope plot shows that the differences in [16 O]/[18 O]
and [16 O]/[17 O] between best estimates of the local ISM and
the disk surrounding VV CrA are not consistent with massdependent isotope fractionation (represented by a line with a
slope of 0.52 in Figure 9). This conclusion is firm to the extent
that it can be assumed that the precursor molecular cloud to
VV CrA was similar to the local ISM. The conclusion is robust despite the uncertainty in the [18 O]/[17 O] ratio of the ISM
because the position of the VV CrA data point on the higher
end of the range in reported [18 O]/[17 O] values (near 4.1) precludes mass-dependent fractionation as a primary control on the
CO isotopologue ratios in all cases. Rather, the trajectories be-
18
tween VV CrA and plausible local ISM isotopic compositions
have slopes of 1 or greater. The slope-1 relationship between
local ISM and CO in the VV CrA disk is reminiscent of the
mass-independent fractionation that characterizes oxygen in the
solar system and is a telltale sign of self-shielding during CO
photolysis.
The ability to distinguish slope-1 (mass-independent) from
slope-0.52 (mass-dependent) trends in Figure 9 is critical for
interpretation of the data and is a key advantage of having data
for all three oxygen isotopologues of CO. The deficit of the
rarer isotopologues, C18 O and C17 O, relative to C16 O in the VV
CrA disk, is on the order of 20% to 40% respectively, which is
significant at the > 2σ level.
The [12 CO]/[13 CO] ratios for both VV CrA and Reipurth
50 are similar, with high-temperature values of 100 ± 10 and
172
SMITH ET AL.
V V CrA
40
16 K
38
Ln(N J/(2J+1))
261 K
36
12
16 K
261 K
16
C O
34
6K
13
16
12
C O
12
C O
C O
176 K
32
15 K
173 K
18
30
0
17
200
400
800
600
Reipurth 50
40
19 K
38
Ln(N J/(2J+1))
152 K
36
12
16
C O
18 K
34
21 K
32
16 K
212 K
13
150 K
116 K
16
C O
30
12
17
C O
0
200
400
12
600
18
C O
800
E/k (K)
Figure 6. Rotational excitation diagrams for CO isotopologues in VV CrA and
Reipurth 50. Simultaneous fits to the two-temperature model and derived high
and low temperatures for each isotopologue are indicated by the red curves and
dashed lines, respectively. Error bars are 1σ propagated from the Gaussian fits.
EJ is the energy of the Jth rotational state above the ground rotational state and
k is the Boltzmann constant. Faded symbols have been excluded from the fits.
110 ± 7, respectively (Table 5). Ratios for both objects are
high compared with the local ISM value of 69 ± 6 (Wilson
1999). In a case where CO self-shielding was almost certainly
responsible for extreme C16 O excesses in the study of the X
Persei molecular cloud by Sheffer et al. (2002), there are no
associated excesses in 12 C, suggesting that carbon and oxygen
isotope effects during CO photolysis can be decoupled. Because
both of our objects have the same [12 C]/[13 C] but one has an
excess in C16 O relative to the ISM, our data also suggest a
decoupling of C and O isotope ratios in CO in these objects.
Self-shielding by CO during photolysis should in principle
result in an excess of 12 CO relative to 13 CO. However, atomic
carbon liberated by CO photolysis is photoionized largely
to C+ , and C+ reacts rapidly to exchange carbon with CO
by the reaction 13 C+ + 12 CO 12 C+ + 13 CO (Warin et al.
1996; van Dishoeck & Black 1988; Langer et al. 1984), hence
driving the 13 C back into CO and diminishing the signature
Vol. 701
of selective photodissociation. In contrast, oxygen liberated
by CO photolysis remains neutral in the photodissociation
regions (PDRs) and is relatively unreactive; oxygen will not
exchange with CO or other molecules prior to sequestration as
H2 O produced either in the gas phase by a reaction channel
involving H+3 (Herbst 2000) or on grain surfaces (Hasegawa
et al. 1992). The decoupling between carbon and oxygen effects
seen here can be explained in the context of photochemical selfshielding but offers as yet no explanation for the anomalously
high [12 CO]/[13 CO] for both VV CrA and Reipurth 50. This
anomaly will be the focus of a later publication.
The cause of the disparities in [16 O]/[18 O] and [16 O]/[17 O]
between the disk surrounding VV CrA and the local ISM
cannot yet be deduced unequivocally, as more data points for
Figure 9 are needed. However, the mass-independent character
of the displacement of this source is most consistent with a
photochemical enhancement of C16 O relative to both C18 O
and C17 O due to self-shielding by the more abundant C16 O
isotopologue.
The absence of a similar effect in Reipurth 50 might be attributable to the different geometry of the absorbing component
of this source. The protostellar envelope in Reipurth 50 is presumably less affected by CO self-shielding than the disk in VV
CrA. If the difference in mass-independent isotope fractionation between disks and protostellar envelopes indeed holds for
more objects, an important question is what is the cause. The
timescale for the isotope-selective photodissociation and subsequent chemistry itself should be short, less than 1000 years, in
the PDR of both types of objects. However, the lines of sight
likely probe more than just the PDR layer so the isotope signature of these other layers is equally relevant. In the disks,
the photodissociation is confined to thin surface regions but the
16
O-depleted atomic oxygen can be mixed vertically down to the
midplane region on timescales of ∼105 years in the outer disk,
where it can then be sequestered into H2 O (Young 2007). Together with the effects of a non-vertical geometry (Section 4.1)
and grain growth by coagulation, this can result in an isotopic
fractionation effect over a much larger column than would be
expected from a thin PDR layer; grain growth can lead to higher
UV flux through the disk and greater photodissociation. In contrast, Reipurth 50, being only in stage I of its evolution, may not
have had sufficient time to acquire a measurable C16 O excess
due to mixing (∼ 105 yr), so that, in the absence of mixing, the
isotope fractionation of the PDR layer contributes little to the
total column.
6. CONCLUSIONS
Our new high-resolution IR absorption data for two YSO
environments, a protoplanetary disk and a protostellar envelope, demonstrate that measurement of both [C16 O]/[C18 O]
and [C16 O]/[C17 O] in CO is possible with sufficient precision to distinguish photochemical effects from mass-dependent
isotope fractionation. The new high-resolution data exhibit a
detectable deficit in C18 O and C17 O in the disk surrounding
VV CrA relative to the local interstellar medium. The CO surrounding Reipurth 50 exhibits no discernible differences in
[C16 O]/[C18 O] and [C16 O]/[C17 O] relative to the local ISM.
One likely explanation for the different oxygen isotopologue
ratios for these objects is that a photochemical deficit in C17 O
and C18 O relative to C16 O proceeds only in a disk geometry and may require 105 years in the environments surrounding YSOs. Future analyses of more objects in different
stages of evolution are now needed to enable differentiation
No. 1, 2009
CO ISOTOPOLOGUE RATIOS IN PROTOSTELLAR DISKS AND ENVELOPES
173
Figure 7. Goodness-of-fit contours for the high temperature lines (J 4) of the four measured isotopologues in VV CrA. The contours show the 68, 95, and 99%
confidence levels. Red stars indicate total column densities obtained from the simultaneous two-temperature model.
Figure 8. Goodness-of-fit contours for the high-temperature lines (J 4) of the four measured isotopologues in Reipurth 50. The contours show the 68, 95, and 99%
confidence levels. Red stars indicate total column densities obtained from the simultaneous two-temperature model.
174
SMITH ET AL.
Vol. 701
Table 4
Measured CO Isotopologue Rotational Temperaturesa
Isotopologue
VV CrA-high T
VV CrA-low T
Reipurth 50-high T
Reipurth 50-low T
261 ± 5
> 100
176 ± 11
173 ± 30
16 ± 1
16 ± 2
6±1
15 ± 5
152 ± 9
212 ± 4
150 ± 3
116 ± 6
19 ± 3
18 ± 1
21 ± 2
16 ± 3
12 C16 O
13 C16 O
12 C18 O
12 C17 O
Note.
a Values were obtained from a simultaneous two-temperature model. All temperatures in K.
Table 5
Measured CO Isotopologue Column Densities and Ratios
Isotopologue Ratios
VV CrA-low T a
VV CrA-high T
7.15 ± 0.17 × 1018
N (12 C16 O)
cm−2
100 ± 10c
690 ± 30
2800 ± 300
4.1 ± 0.4
N (12 C16 O)/N (13 C16 O)
N (12 C16 O)/N (12 C18 O)
N (12 C16 O)/N (12 C17 O)
N (12 C18 O)/N (12 C17 O)
1.40 ± 0.07 × 1018
110 ± 8
1700 ± 250
4200 ± 800
2.5 ± 0.6
Reipurth 50-high T
cm−2
9.03 ± 0.52 × 1018
110 ± 7
490 ± 30
2200 ± 150
4.4 ± 0.2
cm−2
Reipurth 50-low T a
1.64 ± 0.16 × 1018
65 ± 7
340 ± 40
2300 ± 340
6.7 ± 0.8
Local ISMb
cm−2
69 ± 6
557 ± 30
2005 ± 155
3.6 ± 0.2
Notes. Uncertainties reflect the 68% confidence level.
a Low-temperature values are shown for completeness; see text for discussion on their uncertainties.
b (Wilson 1999).
c Assuming T (12 CO) = T (13 CO).
This work was supported in part by a grant from the NASA
Origins program (E.D.Y., M.R.M.) and a grant from the NASA
Astrobiology Institute (UCLA lead team). R.L.S. was supported
by these NASA Origins and Astrobiology grants.
1000
18
17
O/ O = 3.5
O/17O = 4.1
Mass fractionation
Local ISM (Wilson, 1999)
500
REFERENCES
0
RE 50
17
3
17
17
O' = 10 ln( Ri / RLocal ISM)
18
VV CrA
-500
-1000
-1000
-500
18
0
3
18
500
1000
18
O' = 10 ln( Ri / RLocal ISM)
Figure 9. Comparison of oxygen isotope ratios between the interstellar medium
and CO surrounding the two YSOs reported here (stars). Ratios for the YSOs
are derived from the high-temperature CO abundances. The solid and dashed
lines show mass-independent fractionation lines with different assumptions
for the 18 O/17 O ratio, as indicated. The grey line shows the mass-dependent
fractionation line. Ellipses represent 95% confidence limits derived from the
goodness-of-fit contours shown in Figures 7 and 8.
between potential self-shielding in the disk versus the parent
cloud.
We thank the anonymous reviewer for constructive and valuable comments that improved this paper. Support for K.M.P. was
provided by NASA through Hubble Fellowship grant #01201.01
awarded by the Space Telescope Science Institute, which is
operated by the Association of Universities for Research in
Astronomy, Inc., for NASA, under contract NAS 5-26555.
Astrochemistry in Leiden is supported by a Spinoza grant of
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